Characteristics of T-RECS and Michelle data

This page describes the raw data structure for T-ReCS and Michelle
data (imaging and spectroscopy) in their commonly used observing modes,
along with some discussion of different types of noise or instrumental
effects one may see in these images. The basic modes of observation are
discussed on the
Mid-Infrared Observing page. See more information about:

T-ReCS Raw Data

Nearly all T-ReCS observations are taken in chop/nod mode, although a small number of
specialized observations are taken in stare or
stare/nod mode. Given the unusual nature of those observations,
only the chop/nod mode will be discussed in detail here.

The basic feature of all ground-based mid-infrared observation is having
a short exposure time, called the frame time, to avoid saturation
due to the strong thermal emission from the telescope and the atmosphere. The
frame times for T-ReCS filters are generally of the order of 25 ms, although
the values range from about 8 ms for one of the longer wavelength Q-band
filters to 80 ms for a narrow-band filter around 9 microns. The array is
read out frequently and summed in a temporary buffer for each chop position.
These buffers are read out as a saveset. One or more of these savesets
are written out per nod position. In the raw data file there is one extension
per nod position. If there are N savesets per nod, the image written for
each nod has size [320,240,2,N]. The third image dimension selects the chop
position. In nod position A this value is 1 for the on-source (guided) beam
and 2 for the reference (un-guided) beam. In nod position B these are
reversed, the index has a value of 2 for the on-source position and 1 for
the reference position. T-ReCS nods in an ABABABAB pattern so all odd
extensions are for nod A and all even extensions are for nod B.

The header keywords that are relevant to the exposure time and sky
background in T-ReCS images, found in the primary fits header in
extension [0], are the following:

FRMTIME -- the actual integration time of the detector per co-added
frame, in milli-seconds (typical value is about 26 ms)

FRMCOADD -- the number of frames co-added per chop position (a small
number, generally 2 to 5)

CHPSETTL -- the number of frame discarded per chop position to let
the guiding settle after the chop motion on the sky (1 or 2)

CHPCOADD -- the number of chop cycles co-added per saveset

NSAVSETS -- the number of savesets per nod, the size of the fourth
dimension of the raw images (can be 1 to a few)

NNODSETS -- the number of nod positions, either 1 for stare and chop modes or
2 for nod and chop/nod modes

NODSETS -- the number of AB nod cycles for the observation

SAVETIME -- clock time per saveset, in seconds

NODTIME -- clock time per nod position, in seconds

OBJTIME -- the total exposure time in seconds for the
entire observation - on-source, in the guided beam only!

BKAEND, BKASTART, BKAMIN, BKAMAX -- general background values for the
exposure in ADU, usually all very similar values

CHPFREQ -- chop frequency in Hz

T-ReCS raw data files have images of type integer, and as they are a co-add
of a large number of frames the number of ADU per pixel can be quite large.
A background level of several million ADU is quite usual. For example in
taking a normal image of a bright target (actually a standard star)
with the Si-5 11.7 micron filter, the number of savesets per nod was 3, the
frame time was 25.8558006286621 ms (the precision here is excessive,
and it is not clear exactly where it should be rounded off), the number
of frames per chop position was 5, and the number of chop cycles per saveset
was 28. One therefore has about 5 x 28 x 0.025856 seconds on target per
saveset, and 3 of these per nod. This means 10.8595 seconds on target
per nod. With 4 nods total in the exposure the total time on the source
is 43.4381 seconds. The header gives the OBJTIME as 43.4376983642578 seconds,
consistent with the value calculated (and again with excessive numbers of
digits). The background level as given in the header is about 28300 ADU
per frame. 140 of these frames are co-added per saveset, so the magnitude
of the raw signal from the background is about 3.962 million ADU.
The elapsed time is the NODTIME times the number of nods, or 161.626
seconds in this case, so the on-source efficiency is about 27%.

The T-ReCS detector saturates at 65535 ADU per frame, but there is a
"detector clamp level" (given by the RDAEND, RDASTART, RDAMIN, and RDAMAX
header values), so the available dynamic range per frame is roughly
58000 ADU; this background level is roughly 0.6 times the saturation level
and the rest is available for observations of the target. While a very
bright mid-infrared target may saturate the array, this is unusual. Most
mid-infrared targets produce a much smaller signal than does the sky
background.

An individual co-added image for one saveset and one chop position can be
displayed in IRAF with a command such as

cl> display S20070703S0060[1][*,*,1,1] 1 zs+ nsample=76800

Just below is shown the resulting image as seen in the image
display tool. There is an object in the image but it is faint
compared to the sky+telescope background and is not visible in the raw
image.

There are 16 read-out channels for the detector, which causes the
image to have 16 vertical bands each 20 pixels wide as is seen in the
image above. There may be small level shifts between the different
channels, although these are generally small for T-ReCS images. The
other obvious feature seen in the image is a number of bad pixels,
mostly with lower sensitivity, that appear dark on the image.

When the raw data file is "prepared" via the tprepare task the
savesets for each nod are averaged (either with registration or not) and
the difference image is created from the on-source and off-source images.
These are then stored in a image of dimension [320,240,3]. The last index
has values 1, 2, and 3 for the on-source, reference, and difference images
respectively. To display the difference image the command would be

cl> display tS20070703S0060[1][*,*,3] 1 nsample=76800

the result of which is shown just below.

The object is visible in the field. For a single NOD one also sees the
residual telescope radial offset, which appears in the above image as the
bright/dark strips at an angle of about 24 degrees clockwise from the
y axis of the image. The direction of these strips changes with time as the
telescope rotates to keep the image orientation on the sky fixed.

The telescope radiative offset is then removed by summing the nod A
and nod B difference images. Taking the difference images for
extensions 1 and 2 of the image being used as an example here and
adding them together produces an image that is shown below. The
background level is reduced from about 250 ADU to a small value near
zero. One also finds that the channel boundaries are not visible in
the summed image. The magnitude of the residual background depends on
the stability of the atmosphere during the period of observation.
With stable conditions over a pair of nods (roughly 1.5 minutes total
clock time) a good cancellation will result. When that is not the
case one tends to get a relatively uniform positive or negative
baseline in the images along with increased pixel to pixel variations.

When the individual nod difference images are combined to produce a
stacked image pairs of nods are normally averaged together, and then
these averaged images are in turn averaged together. If there is a
problem with an individual nod (because of variable sky conditions,
for instance) then it is necessary to remove that nod and the other
nod of the AB pair from this process.

The negative beams in the stacked images are the two unguided images of
the target. If in the future the guiding can be implemented in both chop
beams using the two peripheral wavefront sensors then for those cases where
two suitable guide stars are available the on-source efficiency can be
doubled.

Instrumental Artifacts

When a bright target is observed there may be artifacts in the raw
data either due to the effect of transferring large amounts of charge
over the array during read-out or due to cross-talk between the
different channels. The most serious of these effects is a loss of
response in a channel if a bright source is in the field. This is
known as the "hammer effect". For a very bright source this will
occur not only in the positive beam but will also be caused by the
unguided images in the other chop position. T-ReCS uses a four-point
sampling algorithm to reduce the hammer effect, although it is not
entirely removed. If a very bright target is observed and the chop
direction/amplitude is such that the object is located in the same
channels on the array in both chop positions then the background level
in those channels will be completely unreliable. Thus for bright
sources it is a good idea to either chop in such a way that the target
is off the detector in the negative beam or to chop along the long
axis of the array to put the negative beams in different channels than
those for the positive beam.

Bright targets also tend to have negative "ghost" images in adjacent
channels. These cannot be removed by changing the chop properties. For a
very bright source the ghost images can become negative streaks along the long
axis of the array.

Below I show an example of an image of a bright compact source. The first
image shows a display of the full intensity range (the chop angle is along the
short axis of the array so the negative beams are not seen). The second image
shows the display with the zs+option in IRAF. In the third image
the range is set to +/- 6000 ADU to show the cross-talk between the channels
more clearly. The IRAF commands for these three cases are:

(The above image shows the entire ds9 display area, including the colour
bar, for reference. For all other images using the "rainbow" colour scheme
the same colour bar applies but the colour bar has not been included. Red
is most intense and pink is least intense.)

In the second image negative regions are seen extending up and down
the channel where the bright core of the object is situated and there are
also negative streaks left and right from the core. A cut across the image
above the core shows the extent of the depression of the sensitivity in the
central channel. The peak of the profile should be smooth and go to a
sharp maximum, but instead one has the noisy pixels within +/-4 pixels of
the position directly above the core--pixel 40 in this plot.

There are other types of noise observed in T-ReCS and Michelle images.
Some of these are specific to different filters. The most common sorts
of such noise show up as either horizontal or vertical stripes in the
difference frames. One such example is shown just below, an image
of a standard star taken in the 10.5 micron [SIV] filter by T-ReCS. The
stripes are seen when one looks at low levels in the image, hence the
use of the zs+ flag in the display. The image was displayed in the
following way:

These sorts of pattern noise in the images are generally fairly
easily removed using the miclean.cl task in the Gemini/IRAF
midir.miutil package, since they are close to being uniform across
the image. One also sometimes encounters similar noise patterns
that are vertical on the array.

Other Observing Modes

In principle T-ReCS observations can be taken in pure chop mode, stare
mode, or pure nod mode. These modes are identified in the header and can
be handled by the basic reduction scripts. In chop-only mode the file
will have one extension with a large number of savesets, but otherwise
be analogous to chop/nod mode. In the more unusual case where chopping is
not used the dimensions of the raw image are altered to reflect the single
chop position. Astronomers who may wish to use T-ReCS in any of these
non-standard modes should contact the T-ReCS instrument scientist through the Helpdesk for
technical evaluation of the suitability and practicality of such
observations before any application for time is submitted.

Spectroscopic Data

T-ReCS spectroscopic observations, like imaging observations, are
taken in chop/nod mode. The same data structure is used for these
observations as is used in imaging. The smaller photon flux per pixel
in this mode means that the hammer effect is only seen for extremely
bright objects. One has the choice of chopping along the slit or
perpendicular to the slit. In the former case one might see the
negative spectra from the second chop position (depending on the angle
and amplitude of the chop), while if chopping off the slit only the positive
spectrum will be seen.

The telescope radiative offset is more difficult to see in the individual
nods for this mode than is the case for imaging observations. The residual
radiative offset is dispersed in wavelength and since typically all parts
of the telescope in the field of view are at similar temperatures the
emission tends to cancel out fairly well between the two chop positions.
Even so one still needs to combine nod pairs to remove this background, and
if a given nod is found to be bad the matching opposite nod should be removed
from the data.

The image below shows a spectrum of a standard star taken with T-ReCS. In
this particular instance only the positive spectrum is seen because the star
is off the slit in the other chop beam. The subsequent figures show a line
cut on the spectrum and the cross-profile of the spectrum. In this image
the wavelength increases from left to right across the array, while the
direction along the slit is along the shorter dimension of the array. The
particular example here shows low resolution N-band spectroscopy. A raw
low resolution Q-band spectrum shows a very different structure due to the
presence of strong atmospheric bands.

On the raw image and on the line cut one sees a lot of
pixel-to-pixel structure. This is not noise as such, it is
fringing in the spectrum. This is a common feature of all such
mid-infrared detectors when used for spectroscopy. The fringing is a
function of the slit width, the spectral resolution, and wavelength.
For some unknown reason the fringing is much worse in the T-ReCS
low-resolution spectroscopy mode compared to the same mode in
Michelle. The fringes can be removed to some degree in the data
reduction, using the msdefringe routine in the Gemini/IRAF
package, but this removal is far from ideal in the higher resolution
spectroscopic mode of T-ReCS and the analogous modes of Michelle
spectroscopy.

The structure in the spectral direction is a combination of the filter
response and the atmospheric transmission. In the middle of the spectrum
around x=135 is a strong ozone band. Around x=90 there is another strong
dip in the spectrum which is due to a feature in the N-band filter response
(and which is unique to T-ReCS). At left and right, x less than 45 or
larger than 300, there is no spectral response due to cut-offs from the filter.
The plot just below shows a comparison of a raw T-ReCS low resolution N-band
spectrum to the expected spectrum of the star after transmission through the
atmosphere are this resolution. The ozone feature is obvious in the
transmitted spectrum plots, as are strong lines of water and other molecules
at wavelengths outside the filter passband. The raw T-ReCS spectrum is
shown with the y-axis scale at right, while the other three curves have
y-axis scaling as at left in this figure.

The ozone band is somewhat variable with time and direction, and since
the band is so strong using a telluric calibration source to remove the band
is essential in the data reduction for such N-band spectra.

Chop-only Mode: Spectroscopic Acquisitions

In T-ReCS, as in Michelle, spectroscopic acquisition is usually done in
chop-only mode. The images are of dimension [320,240,2,N] where N is some
number of savesets.

The T-ReCS PSF Structure

Observations in the mid-infrared have high Strehl ratios (typical
values are 35% in the short wavelength end of N-band, 65% in the long
wavelength end of N-band, and 85% in Q-band; in very good seeing the
estimated Strehl in the Si-5 11.7 micron filter can reach 85%). Under
most optical seeing conditions the point spread function has full
width at half maximum (FWHM) which is close to the theoretical
resolution of the telescope. For Gemini at N-band with good seeing
the FWHM in the Si-5 11.7 micron filter is usually near 0.35 to 0.37
arc-seconds (see the Observing
Conditions Constraints page). In good conditions the FWHM tends
to be roughly independent of the airmass -- as opposed to the case in
the optical where the FWHM generally scales as airmass to the 0.6
power, in which case for stable seeing conditions the FWHM at 2
airmasses is 1.5 times what it is at zenith. At N-band one can have
stable 0.35 arc-second FWHM seeing from the zenith down to two
airmasses with very little change. Ideally the FWHM will be
proportional to the wavelength through N-band and Q-band. In practice
we see a somewhat flatter variation in seeing with the filter central
wavelength (i.e., the IQ does not improve as quickly towards shorter
wavelengths as would be expected from a simple scaling with
wavelength).

Only under very poor optical seeing do we begin to see signifaicant
atmospheric effects on the PSF at N-band and Q-band, whereupon the FWHM
can become as large as 0.7 to 1.0 arc-seconds in the N-band filters
(although this is relatively rare). In such cases where the bulk of
the seeing is due to the atmospheric turbulence the FWHM begins to scale
with airmass in the same way as the optical seeing normally does.

It is found that under the best seeing conditions at N-band or
Q-band the Strehl ratio increases while the FWHM of the PSF does not
change much. When the seeing gets somewhat worse the first indication
is a loss of power in the PSF peak without significant change in the
FWHM, followed by a broadening of the peak and a change of the shape
of the PSF to follow a pure Moffat profile similar to what is normally
observed with the (optical) GMOS instrument. Under better seeing
conditions the PSF core is slightly better fit with a Gaussian profile
than with a Moffat profile, although the difference is usually subtle.

Under good seeing the T-ReCS PSF shows several diffraction rings around
the central core. Generally the trefoil is prominent around the peak and
one or two outer diffraction rings can be seen in the image when a standard
star is observed under good seeing. An example of this is shown below. The
first image is a linear display of an image of alpha CMa zoomed in to the
star. It shows the trefoil around the core in a fairly typical manner.
The second image shows a the full frame of another observation of the same
star on a different date with very good seeing (note that the orientation
of the trefoil has changed...), where the image maximum is reduced by a
factor of 5 from the peak value to highlight the low-level structure in
the PSF. The third image is the same as the second image only zoomed in
somewhat to show the structure better.

The three-fold symmetry in the inner diffraction ring, trefoil, is typical
of mid-infrared images on Gemini. When the PWFS2 is used for guiding the
trefoil is not corrected for. In the less usual case where PWFS1 is used
to guide the wavefront sensor sends trefoil corrections to the telescope and
this trefoil structure is removed. The trefoil tends to be stronger in the
early T-ReCS data (to 2005) but has been less prominent in more images
since 2007. As most of the T-ReCS examples here are drawn from older images
the trefoil is often quite obvious.

For this image there is a filter ghost north-north-west of the core at
about 2 o'clock just outside the second diffraction ring. Although it is
only faintly seen by eye in the above images, there is also a third
diffraction ring in the image. Making a radial plot of the PSF allows one
to see the diffraction rings as is seen below.

The core FWHM is 4.12 pixels = 0.37 arc-seconds. The trefoil produces the
first bump on the profile outside the core at radius about 6 pixels. The
next diffraction ring is at radius 13 pixels and there is another at radius
20.5 pixels. The more diffuse peak at radius 17 pixels corresponds to the
filter ghost. There may be another much fainter diffraction ring at radius
32 pixels, although this is marginal at best.

The way the radial profile looks using the default IRAF rimexam
parameters is shown just below. In such a plot one sees the core and the
trefoil components.

An example of the PSF and the radial profile under poorer seeing
conditions (definite optical
IQ=85% conditions) is shown below. The FWHM is now about 15%
larger than it would be in good seeing, and the radial profile follows
the Moffat profile fairly well over the range of the plot. One no
longer sees the trefoil due to the poorer intrinsic seeing conditions.
There is also some odd low level structure in the PSF and the
ellipticity is higher than it usually is in good seeing (e=0.10 here
compared to e=0.05 for the alpha CMa image above in good seeing). As
long as the PSF is stable over the time of the observation images of
extended targets can be deconvolved even with a poorer PSF like this
to get a good reconstructed image.

For comparison consider the following two PSF profiles. The
first case is for the same object as in the above example, under
average to a bit worse than average seeing in N-band: the FWHM in
the 11.7 micron Si-5 filter is about 0.40 arc-seconds which
is not as good as is often observed (values down to 0.37 arc-seconds
can be obtained in better seeing conditions), but the estimated Strehl
is 0.54 and one can see the prominent trefoil around the core of the image.
The radial profile is shown after the image below. The trefoil causes the
radial profile to be above the Moffat fit from about 4 to 9 pixels radius.
This is the general type of profile expected under IQ=70% conditions in the
mid-infrared.

Under poorer seeing conditions, IQ=85% in the mid-IR, and with the same
filter the PSF and profile look more like the example given below. The
FWHM has increased to 0.48 arc-seconds, while the estimated Strehl has
decreased to 0.28, only about half the value in the case just above. When
the optical seeing is extremely poor one may get FWHM values of the order
of 1 arc-second in N-band. As noted above, in such poor seeing conditions
the PSF scales with airmass in a way similar to what is normally the case
in the near-infrared and the optical.

Michelle Raw Data

Michelle operates very similarly to T-ReCS, so most of the features
of the raw and stacked data files are very similar. One significant
difference between the two instruments is that Michelle data are saved
at the end of one full nod period, roughly 40 seconds in most cases
(no savesets are used). Michelle imaging and low resolution
spectroscopic observations are taken in chop/nod mode as is also the
case for T-ReCS. However the higher resolution spectroscopic modes of
Michelle are taken in pure nod mode without chopping, which is quite
different than the case for T-ReCS.

In chop-nod mode the Michelle raw data files have images for the two chop
beams and a difference image right at the start. Thus the image format of
each extension of a raw Michelle file is [320,240,3,1]. The last dimension is
degenerate and only present so that T-ReCS and Michelle raw data images have
the same dimensionality. In the raw data the third dimension has values
1 for the difference image, 2 for chop position A and 3 for chop position B.
The difference is always chop position A minus chop position B, so the
difference changes sign from the A nods to the B nods.

The header keywords concerning the exposure time for Michelle are the
following:

EXPOSURE -- the time per frame in seconds, usually 0.02 to 0.09 for
imaging

NUMEXPOS -- the number of frames co-added in per nod

NUMEXT -- the number of extensions in the image which is also the
number of nods used, and should be a multiple of 4 for normal ABBA observations

OBSTIME -- the total time elapsed for the observation

CHOPFREQ -- chop frequency in Hz

The time on-source in the guided beam is given by the frame
time multiplied by the number of exposures. The total time for the
exposure is also given as the OBSTIME header keyword, so the on-source
efficiency can be calculated if desired from these values. An example
image with the 11.7 micron Si-5 filter has an EXPOSURE value of 0.09
seconds, a NUMEXPOS value of 144, a NUMEXT value of 8, and an OBSTIME
of 365.999221801758 seconds. The total time on source is 0.09 x 144 x
8 seconds = 103.68 seconds. The on-source efficiency is therefore
about 28%.

One distinct difference between Michelle data and T-ReCS data is that
in Michelle the frames within a nod are averaged rather than co-added.
The data type of the raw data files is real numbers rather than integers.
The raw background per chop position is seen in the raw images directly.
For Michelle saturation occurs at around 52000 ADU so the raw sky background
should be much lower than this (roughly half full well in average weather
conditions) to preserve a good part of the detector dynamic range for
detection of the object.

An on-source raw image is shown below. One thing that is seen right off
is that there is a vignetted region in the left corners, which is true for all
images taken by Michelle. The same 16 channel structure is seen here as in
T-ReCS raw frames, but the channel to channel offsets are somewhat more obvious
than is the case for T-ReCS. Michelle does not use four-point sampling when
reading the array out. The following line cut across the raw image shows
the channels quite clearly. Each group of four channels from left to right
has a separate read-out channel and thus tends to be offset from the other
groups. The channel offsets can be removed by taking a bias frame (although
that is not normally done except for spectroscopic observations) or by
software. The miclean.cl task generally is able to remove the channel
offsets from Michelle images. In most cases the channel offsets are much
reduced in the difference images.

The hammer effect is typically much more of a problem in Michelle images
than in T-ReCS images, again due to not using four-point sampling of the
array during read-out. A rather bad example of the hammer effect is shown
below where an image was taken of an extremely bright mid-infrared source,
IK Tau. The hammer effect up the columns in the middle are at the 3% level
compared to the peak on the target. The second picture shows the same image
with the zs+ option, whereupon the hammer effect from the negative beam
at right is clearly seen. (There is no hammer effect from the other negative
beam because the source is at the top of the columns.)

In the difference image one does not see the vignetted regions in the left
corners because the general background level is near zero. However if the
background becomes high enough to saturate the raw images the shape of these
vignetted corners tends to be seen in the Michelle difference images, and
hence in the stacked difference image.

A raw Michelle low resolution N-band spectral frame is shown below.
In this case the star is chopped along the slit, and the two negative
spectra are seen. For reasons that are not yet understood the lower
of the two negative beams is generally much weaker than the upper one,
whereas normally one would expect the two to be of the same strength
and each about half the strength of the positive spectrum.

A cut across the spectral image shows a fairly similar ozone band structure
as is seen in the T-ReCS raw spectral cut above--note however the lack of
fringing. As Michelle uses a long-pass filter which cuts on at about 7
microns when doing low resolution N-band spectroscopy, when conditions are
dry there is some spectral signal over most of the array. The structure
in the cut on the left-hand side reflects water bands in the atmosphere. On
the right-hand side of the array there is only very weak signal. While the
filter passes radiation beyond 13 microns, there is strong atmospheric
emission at these wavelengths in addition to radiation from near 7 microns
coming in at second order. As a result, Michelle is unable to detect
the spectrum much beyond 13.5 microns, although it does have slightly better
long wavelength coverage than T-ReCS does.

Polarimetry Mode

The polarimetry mode of Michelle produces a different type of raw data
file. Each extension contains the images for one nod position, but within
the extension there are images for several waveplate positions. The
current mode of observation is to have eight waveplate positions observed
per nod -- at waveplate angles of 0, 45, 45, 0, 22.5, 67.5, 67.5, and
22.5 degrees. The images from each waveplate position have the same
structure as in normal imaging mode, [320,240,3], and when written out
to the raw data file the eight images are combined into an image of
dimension [320,240,3,8] where the last dimension is the index of the
waveplate positions. The purpose of having eight waveplate positions
per nod in this manner is to take out any linear drifts in the system,
being directly analogous to ABBA nodding. However the efficiency is low
due to the frequent changes of the waveplate position per nod (the nods
are longer than in normal imaging mode, but only by about a factor of 2
so the time per waveplate position is quite short).

Once a raw Michelle polarimetry data file is processed with the
mprepare task in the GEMINI/IRAF package the images for each
waveplate position are split out into separate extensions.

Noise in Michelle Images

The same types of pattern or other noise are seen in Michelle images
as are seen in T-ReCS images. The channel to channel offsets are usually
more obvious in Michelle stacked images than in T-ReCS ones, but the
miclean.cl routine removes these offsets quite well so this is not
usually an issue. The various types of stripe noise seen in T-ReCS
images are also seen in Michelle images, although they are rarer in general
(possibly because T-ReCS has a much wider selection of filters than does
Michelle, as many of these types of noise are filter-specific).

The figures below show in turn (a) a stacked image of a fairly bright
standard star displaying the full intensity range, (b) the same image
showing only the range of pixels near the median, in which one sees small
channel offsets which we call "staircasing", (c) the same image after the
channel to channel offsets are removed with miclean.cl, wherein
there is some banding across the image from left to right, and (d) the
same image again after the horizontal noise is removed with miclean.cl.
One can see that the systematic background is removed rather well by the IRAF
script. The various commands used in IRAF for this are:

Nod-only Mode: Higher Resolution Spectroscopy

For the Michelle low-N and Low-Q spectroscopic modes the observations are
taken in the same chop/nod mode as imaging observations, although with a
rather lower chop frequency of about 1 Hz rather than 3 to 4 Hz as for
imaging. However all the higher resolution Michelle spectroscopy modes --
medN1, medN2, and echelle -- are taken in the same way as near-infrared
spectral observations: nodding along the slit to remove the background
emission, rather than chopping and nodding.

The individual file extensions in the raw data have dimension [320,240,1,1]
so the last two dimensions are degenerate. This is done to keep the same
number of dimensions as in the chop/nod mode case (which in turn has the extra
dimension to match the T-ReCS raw data files). Just below are shown a raw
image of a medN1 spectrum with central wavelength of 11.5 microns. The star
is much fainter than the background. On the raw image one sees two things:
the atmospheric emission lines which are almost straight along the y axis of
the image, and the fringes which in this case are tilted slightly clockwise
from the y axis. There is an un-illuminated region at the top of the array,
this is normal in the higher resolution spectral modes of Michelle where a
steeper grating tilt is used.

The next image shows the simple difference of the first two nod positions,
where the positive and negative stellar spectra are visible. There are still
some residual sky lines visible due to small variations in the conditions
during the two nods. These tend to disappear when more nods are stacked
up, in part because any linear time trends in the sky line strengths
are taken out by the ABBA nod pattern. After that the stacked image is shown.
In this case there is very little in the way of residual sky emission. The
case where residual lines are more likely is where lines are saturated in the
exposures. The region of N-band in this spectrum is mostly clear of very
strong lines, so this effect is not particularly evident here. Where these
lines are present no signal is received anyway, so it does not matter that
there is some residual signal (the residual signal is due to variations in
what part of the core of the line is saturated, and can leave either
positive or negative residuals depending on how conditions change during
the observation).

The magnitude of the sky emission and the fringing is shown in the
lines cuts below. The first one is a line-cut at the source position in
the raw image above. The following plot shows the same line cut in the
stacked difference image. Note the change in total intensity and that the
fringes are not removed in this process. All of, or nearly all of, the
high frequency structure seen in this spectral cut is due to the fringing.
Unfortunately the fringing signal does not appear to have a simple Fourier
signature that can be removed easily from the spectrum.

Finally, the plot below shows a profile across the spectrum along the
slit direction. This is fairly typical of Michelle spectra.

Chop-only Mode: Spectroscopic Acquisitions

In most cases spectroscopic acquisitions for Michelle are done in
chop-only mode, which is treated like a single nod of a normal chop/nod
observation. Thus the background radiative offset of the telescope is
present in these images. The only other matter of note for these observations
is that the orientation of the field is rotated 90 degrees from the standard
orientation, so north is to the left and east is down on the array.

Stare Mode: Flats and Biases

When spectral observations are taken Gemini provides a flat image and
a bias image to aid in the spectral data reduction. These are taken in
simple stare mode. The image dimensions are the same as for nod mode, since
it is treated like a single nod observation. The flat is made by having
Michelle look at an internal surface which is at nearly ambient temperature,
so to a fair approximation the spectrum is that of a blackbody at roughly
0 degrees Celsius. In low resolution N-band or Q-band modes the brightness
changes a fair amount over the array. In all the higher resolution modes
there is not much spectral structure in the flats. The third image is the
difference between the flat and the bias image. This should in principle
be the pure spectrum of the surface. It does have some structure still
present, presumably due to the material in the surface being looked at.

A line cut through the flat is shown below. There is some structure
that is seen on the cut, aside from the overall spectral rise and
turn-over. The second line cut shows the residual structure after the
bias removal. A line cut from the bias is also shown as the last plot,
wherein one sees the staircasing between the channels.

A low-N flat image

A bias image

A low-N bias-subtracted flat image

The Michelle PSF Structure

In general terms all the comments about the T-ReCS PSFs
apply to the Michelle PSF as well. The Michelle pixels are slightly larger as
projected on the sky than the T-ReCS ones (0.1005 arc-seconds across compared to
0.08976 arc-seconds across) which makes small differences in sampling but otherwise
the PSFs nominally have the same characteristics as for T-ReCS.

However there was a problem with the Michelle PSFs that is not
observed with T-ReCS and which turned out to be a telescope guiding
problem rather than a problem inside Michelle itself: the PSFs had a
tendency to be elongated in the chop direction. This problem was present in Michelle images taken prior to June 2008, although the magnitude of the effect appeared to be worse after 2006 than it had been prior to that time. This problem with the guiding was fixed by engineering work in 2008, and it should not appear in Michelle data taken after May 2008. The discussion here is therefore for "historical" reference.

These elongations are
colloquially known as "chop tails". They cause the typical
ellipticity of the Michelle PSFs in good seeing conditions to be
e=0.15 to e=0.20 as measured in IRAF, compared to typical values
e=0.05 in cases where the chop tails are not seen in Michelle
observations. Examination of the individual nod images when the chop
tails are present show that the elongation is in the negative chop
direction for each nod position, so that when the A and B nods are
combined the result is elongation on both sides of the core of the
PSF. The cause of these chop tails is not yet understood, and while
the chop tails are not always present it appears that the frequency of
occurrence of the chop tails has increased with time from 2004 through
2007. Early on only a small fraction of images were affected by the chop direction
elongation, whereas by semester 2007B almost all observations have
this problem.

An example of the PSF one obtains with the chop tails problem is shown
below. This is from an image taken under good seeing conditions on September
14, 2007. One sees the core, the usual trefoil and first diffraction ring,
and a second fainter diffraction ring (the figures also label another
"diffraction ring" between the trefoil and the first complete outer
diffraction ring, but this may be an extension of the trefoil due to the
chop tails). The core is seen to be somewhat elongated, close to along the
y axis of the plot which is the chop direction in this observation. The
nominal angle from the PSF fit is -9 degrees E of N, slightly off of the
chop direction (which is also typical of this problem). The images below
show the PSF in a linear plot, in a logarithmic plot, then with the features
marked. The last figure shows a radial profile for the star, and the blue
arrows mark the diffraction rings. As this source is 3.5 times fainter
than alpha CMa (42.6 Jy compared to 143.1 Jy in the IRAS [12] filter) the
peak to background ratio is smaller than in the example radial profiles for
T-ReCS above. It is a good idea to use somewhat fainter stars with
Michelle because the hammer effect is worse than with T-ReCS.

The ellipticity of the image core is e=0.10 from IRAF. This is at
least twice what it should be under these seeing conditions. In worse
cases the ellipticity can be 0.2 or sometimes higher due to the chop
tails. There is some evidence that the chop tails are mimimized when
chopping east-west and are stronger when chopping
in other directions, but this is yet to be rigorously tested.

The flux level of the outer diffraction ring in the PSF images is about
0.00015 times the raw background sky flux. Nonetheless it is visible in
one ABBA nod cycle. The peak of the stellar PSF for this object of
brightness 28.38 Jy in this filter is still 8 or 9 times less than
the background radiation in a raw frame. Calibration of this image implies
a sky brightness of roughly 1700 Jy/square arc-second in the Si-5 filter,
and a brightness of 0.24 Jy/square arc-second in the outer diffraction ring.

Characteristics of Saturated Images (T-ReCS or Michelle)

If mid-infrared observations are taken under variable cloud
conditions, or if some object is blocking the field of view, the raw
Michelle or T-ReCS images may be saturated. While the observer at the
telescope would be expected to see that the image is saturated, and
usually to flag it as "fail", in this section a discussion is given of
how to detect saturation of the images. If only a few nods in a long
observation are saturated the observation may be set to "pass", either
because the saturated nods were not seen at the telescope or because
the conditions are consistent with what is requested (CC70 or worse
conditions may well result in some saturated data, especially in
imaging mode). Normally an observation would not be passed if a
wavefront sensor, or some other unusual part of the telescope, is in
the beam, but there are rare exceptions to this as well.

The example image used to illustrate the effects of saturation here is one
where the only PWFS2 guide star available is only 3.67 arc-minutes from the
target position. In the OT one finds that the outline of PWFS2 on the image
display overlaps on the image of the field (see below). The outline in the OT
image widget is quite accurate, although the affected area is slightly larger
than would strictly be expected from the shadowed region shown.

In this instance since no other guide star was available and since the
target object was expected to be point-like the sky rotation of the image
was selected to minimize the shadowing of the field by the wavefront sensor.
In the thermal infrared the probe not only shadows the field but in fact
shines such that it saturated the detector along the top edge and in the
upper right corner. A example of a raw NOD frame is shown below.

A line cut across the image shows the unusual background produced by the
radiation from PWFS2 seen by Michelle; it is not limited to the regions
indicated in the OT but extends to the middle of the array. The background
in this filter would normally be about 35000 ADU.

The saturation level for Michelle is about 56000 ADU. This brightness
level is reached in the upper right corner of the raw nod images, and
line cuts become flat in this region.

When the difference images are taken and are stacked with the
mireduce Gemini/IRAF task the direct evidence of the saturation is
lost. However one can still detect that the upper right corner of the
resulting image has unusual noise properties when one plots the resulting
image with the "zs+" option. This is shown below. (The target of this
observation is too faint to see in this image. One sees only noise across
the frame.)

The sign of saturation is the suppression of the pixel to pixel
variations in the frame. One also sees a faint stripe pattern at
an angle of about 45 degrees E of N in the upper right corner.
A cut through this region shows the reduction in pixel to pixel fluctuations
in the stacked image where the probe was in the field of view. See the line
cut below, where the pixel to pixel fluctuations are reduced past about pixel
180 in the image.

If a Michelle or T-ReCS observation is taken in cloudy conditions it may
be that some of the individual nods saturate due to clouds passing through
the field. Such cases are rare, since we normally do not observe with
the mid-IR instruments in worse than light CC=70% conditions. However when this
happens what is usually seen in the stacked image is a reduction in the
dynamic range. One either sees a flattening of the profile across the
object or a suppression of the pixel to pixel fluctuations in the resulting
image, depending upon the severity of the saturation in the nods. The only
way to be certain that nods are not saturated is to examine the individual
nods of an observation and remove those that are adversely affected by the
clouds.